Contents

Astronomical interferometers come in two types—direct detection and heterodyne. These differ only in the way that the signal is transmitted. Aperture synthesis can be used to computationally simulate a large telescope aperture from either type of interferometer.

In the near future other arrays are expected to release their first interferometric images, including the ISI, VLTI, the CHARA array and the MRO interferometers.

At the beginning of the 21st Century, the VLTI and Keck Interferometer large-telescope arrays came into operation, and the first interferometric measurements of the brightest few extra-galactic targets were performed.

A simple two-element optical interferometer. Light from two small telescopes (shown as lenses) is combined using beam splitters at detectors 1, 2, 3 and 4. The elements create a 1/4 wave delay in the light, allowing the phase and amplitude of the interference visibility to be measured, thus giving information about the shape of the light source.

A single large telescope with an aperture mask over it (labelled Mask), only allowing light through two small holes. The optical paths to detectors 1, 2, 3 and 4 are the same as in the left-hand figure, so this setup will give identical results. By moving the holes in the aperture mask and taking repeated measurements, images can be created using aperture synthesis, which would have the same quality as would have been given by the right-hand telescope without the aperture mask. In an analogous way, the same image quality can be achieved by moving the small telescopes around in the left-hand figure – this is the basis of aperture synthesis, using widely separated small telescopes to simulate a giant telescope.

One of the first astronomical interferometers was built on the Mount Wilson Observatory's reflector telescope in order to measure the diameters of stars. This method was extended to measurements using separated telescopes by Johnson, Betz and Towns (1974) in the infrared and by Labeyrie (1975) in the visible. The red giant star Betelgeuse was among the first to have its diameter
determined in this way. In the late 1970s improvements in computer processing allowed for the first "fringe-tracking" interferometer, which operates fast enough to follow the blurring effects of astronomical seeing, leading to the Mk I, II and III series of interferometers. Similar techniques have now been applied at other astronomical telescope arrays, such as the Keck Interferometer and the Palomar Testbed Interferometer.

A detailed description of the development of astronomical optical interferometry can be found here. Impressive results were obtained in the 1990s, with the Mark III measuring diameters of hundreds of stars and many accurate stellar positions, COAST and NPOI producing many very high resolution images, and ISI measuring stars in the mid-infrared for the first time. Additional results included direct measurements of the sizes of and distances to Cepheid variable stars, and young stellar objects.

Interferometers are seen by most astronomers as very specialized instruments, as they are capable of a very limited range of observations. It is often said that an interferometer achieves the effect of a telescope the size of the distance between the apertures; this is only true in the limited sense of angular resolution. The combined effects of limited aperture area and atmospheric turbulence generally limit interferometers to observations of comparatively bright stars and active galactic nuclei. However, they have proven useful for making very high precision measurements of simple stellar parameters such as size and position (astrometry) and for imaging the nearest giant stars. For details of individual instruments, see the list of astronomical interferometers at visible and infrared wavelengths.

Radio wavelengths are much longer than optical wavelengths, and the observing stations in radio astronomical interferometers are correspondingly further apart. The very large distances do not always allow any usable transmission of radio waves received at the telescopes to some central interferometry point. For this reason many telescopes instead record the radio waves onto a storage medium. The recordings are then transferred to a central correlator station where the waves are interfered. Historically the recordings were analog and were made on magnetic tapes. This was quickly superseded by the current method of digitizing the radio waves, and then either storing the data onto computer hard disks for later shipping, or streaming the digital data directly over a telecommunications network e.g. over the Internet to the correlator station. Radio arrays with a very broad bandwidth, and also some older arrays, transmit the data in analogue form either electrically or through fibre-optics. A similar approach is also used at some submillimetre and infrared interferometers, such as the Infrared Spatial Interferometer. Some early radio interferometers operated as intensity interferometers, transmitting measurements of the signal intensity over electrical cables to a central correlator. A similar approach was used at optical wavelengths by the Narrabri Stellar Intensity Interferometer to make the first large-scale survey of stellar diameters in the 1970s.

At the correlator station, the actual interferometer is synthesized by processing the digital signals using correlator hardware or software. Common correlator types are the FX and XF correlators. The current trend is towards software correlators running on consumer PCs or similar enterprise hardware. There also exists some amateur radio astronomy digital interferometers, such as the ALLBIN of the European Radio Astronomy Club.

As most radio astronomy interferometers are digital they do have some shortcomings due to the sampling and quantization effects as well as the need for much more computing power when compared to analog correlation. The output of both a digital and analog correlator can be used to computationally synthesize the interferometer aperture in the same way as with direct detection interferometers (see above).

1.
Astronomical interferometer
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The advantage of this technique is that it can theoretically produce images with the angular resolution of a huge telescope with an aperture equal to the separation between the component telescopes. The main drawback is that it does not collect as much light as the complete instruments mirror, thus it is mainly useful for fine resolution of more luminous astronomical objects, such as close binary stars. Another drawback is that the angular size of a detectable emission source is limited by the minimum gap between detectors in the collector array. Interferometry is most widely used in astronomy, in which signals from separate radio telescopes are combined. A mathematical signal processing technique called aperture synthesis is used to combine the signals to create high-resolution images. Practical infrared and optical astronomical interferometers have only recently been developed, at optical wavelengths, aperture synthesis allows the atmospheric seeing resolution limit to be overcome, allowing the angular resolution to reach the diffraction limit of the optics. Astronomical interferometers can produce higher resolution images than any other type of telescope. At radio wavelengths, image resolutions of a few micro-arcseconds have been obtained, one simple layout of an astronomical interferometer is a parabolic arrangement of mirror pieces, giving a partially complete reflecting telescope but with a sparse or dilute aperture. Instead, most existing arrays use a geometry, and Labeyries hypertelescope will use a spherical geometry. One of the first uses of optical interferometry was applied by the Michelson stellar interferometer on the Mount Wilson Observatorys reflector telescope to measure the diameters of stars, the red giant star Betelgeuse was the first to have its diameter determined in this way on December 13,1920. In the 1940s radio interferometry was used to perform the first high resolution radio astronomy observations, optical/infrared interferometry was extended to measurements using separated telescopes by Johnson, Betz and Townes in the infrared and by Labeyrie in the visible. Similar techniques have now been applied at other astronomical telescope arrays, including the Keck Interferometer, software packages such as BSMEM or MIRA are used to convert the measured visibility amplitudes and closure phases into astronomical images. If completed, the MRO Interferometer with up to ten movable telescopes will produce among the first higher fidelity images from a long baseline interferometer, astronomical interferometry is principally conducted using Michelson interferometers. The principal operational interferometric observatories which use this type of instrumentation include VLTI, NPOI, engineers at the European Southern Observatory ESO designed the Very Large Telescope VLT so that it can also be used as an interferometer. Along with the four 8. 2-metre unit telescopes, four mobile 1. 8-metre auxiliary telescopes were included in the overall VLT concept to form the Very Large Telescope Interferometer. The ATs can move between 30 different stations, and at present, the telescopes can form groups of two or three for interferometry, when using interferometry, a complex system of mirrors brings the light from the different telescopes to the astronomical instruments where it is combined and processed. This is technically demanding as the light paths must be equal to within 1/1000 mm over distances of a few hundred metres. For the Unit Telescopes, this gives an equivalent mirror diameter of up to 130 metres, and this is up to 25 times better than the resolution of a single VLT unit telescope

2.
Visible-light astronomy
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Visible-light astronomy encompasses a wide variety of observations via telescopes that are sensitive in the range of visible light. An example of spectroscopy is the study of lines to understand of what kind of matter light is going through. Visible astronomy also includes looking up at night and he later made improved versions with up to about 30x magnification. With a Galilean telescope the observer could see magnified, upright images on the earth—it was what is known as a terrestrial telescope or a spyglass. He could also use it to observe the sky, for a time he was one of those who could construct telescopes good enough for that purpose. On 25 August 1609, he demonstrated one of his telescopes, with a magnification of about 8 or 9. His telescopes were also a sideline for Galileo selling them to merchants who found them useful both at sea and as items of trade. He published his initial telescopic astronomical observations in March 1610 in a treatise entitled Sidereus Nuncius. The visibility of objects in the night sky is affected by light pollution. The presence of the Moon in the sky has historically hindered astronomical observation by increasing the amount of ambient lighting. With the advent of light sources, however, light pollution has been a growing problem for viewing the night sky. List of largest optical reflecting telescopes

3.
Interferometry
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Interferometry is a family of techniques in which waves, usually electromagnetic waves, are superimposed causing the phenomenon of interference in order to extract information. Interferometers are widely used in science and industry for the measurement of small displacements, refractive index changes, in an interferometer, light from a single source is split into two beams that travel different optical paths, then combined again to produce interference. The resulting interference fringes give information about the difference in path length, waves which are not completely in phase nor completely out of phase will have an intermediate intensity pattern, which can be used to determine their relative phase difference. Most interferometers use light or some form of electromagnetic wave. Typically a single incoming beam of coherent light will be split into two beams by a beam splitter. Each of these beams travels a different route, called a path, the path difference, the difference in the distance traveled by each beam, creates a phase difference between them. It is this phase difference that creates the interference pattern between the initially identical waves. If a single beam has been split along two paths, then the difference is diagnostic of anything that changes the phase along the paths. This could be a change in the path length itself or a change in the refractive index along the path. As seen in Fig. 2a and 2b, the observer has a view of mirror M1 seen through the beam splitter. The fringes can be interpreted as the result of interference between light coming from the two virtual images S1 and S2 of the original source S, the characteristics of the interference pattern depend on the nature of the light source and the precise orientation of the mirrors and beam splitter. In Fig. 2a, the elements are oriented so that S1 and S2 are in line with the observer. Use of white light will result in a pattern of colored fringes, the central fringe representing equal path length may be light or dark depending on the number of phase inversions experienced by the two beams as they traverse the optical system. Interferometers and interferometric techniques may be categorized by a variety of criteria, In homodyne detection, the phase difference between the two beams results in a change in the intensity of the light on the detector. The resulting intensity of the light after mixing of two beams is measured, or the pattern of interference fringes is viewed or recorded. Most of the interferometers discussed in this article fall into this category, the heterodyne technique is used for shifting an input signal into a new frequency range as well as amplifying a weak input signal. A weak input signal of frequency f1 is mixed with a reference frequency f2 from a local oscillator. The nonlinear combination of the input signals creates two new signals, one at the sum f1 + f2 of the two frequencies, and the other at the difference f1 − f2 and these new frequencies are called heterodynes

4.
Telescope
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A telescope is an optical instrument that aids in the observation of remote objects by collecting electromagnetic radiation. The first known practical telescopes were invented in the Netherlands at the beginning of the 1600s and they found use in both terrestrial applications and astronomy. Within a few decades, the telescope was invented, which used mirrors to collect. In the 20th century many new types of telescopes were invented, including radio telescopes in the 1930s, the word telescope now refers to a wide range of instruments capable of detecting different regions of the electromagnetic spectrum, and in some cases other types of detectors. The word telescope was coined in 1611 by the Greek mathematician Giovanni Demisiani for one of Galileo Galileis instruments presented at a banquet at the Accademia dei Lincei, in the Starry Messenger, Galileo had used the term perspicillum. The earliest recorded working telescopes were the telescopes that appeared in the Netherlands in 1608. Their development is credited to three individuals, Hans Lippershey and Zacharias Janssen, who were spectacle makers in Middelburg, Galileo heard about the Dutch telescope in June 1609, built his own within a month, and improved upon the design in the following year. Also in 1609, Thomas Harriot became the first person known to point a telescope skyward in order to make observations of a celestial object. The idea that the objective, or light-gathering element, could be a mirror instead of a lens was being investigated soon after the invention of the refracting telescope. The potential advantages of using parabolic mirrors—reduction of spherical aberration and no chromatic aberration—led to many proposed designs, in 1668, Isaac Newton built the first practical reflecting telescope, of a design which now bears his name, the Newtonian reflector. The invention of the lens in 1733 partially corrected color aberrations present in the simple lens and enabled the construction of shorter. The largest reflecting telescopes currently have objectives larger than 10 m, the 20th century also saw the development of telescopes that worked in a wide range of wavelengths from radio to gamma-rays. The first purpose built radio telescope went into operation in 1937, since then, a tremendous variety of complex astronomical instruments have been developed. The name telescope covers a range of instruments. Most detect electromagnetic radiation, but there are differences in how astronomers must go about collecting light in different frequency bands. The near-infrared can be collected much like light, however in the far-infrared and submillimetre range. For example, the James Clerk Maxwell Telescope observes from wavelengths from 3 μm to 2000 μm, on the other hand, the Spitzer Space Telescope, observing from about 3 μm to 180 μm uses a mirror. Also using reflecting optics, the Hubble Space Telescope with Wide Field Camera 3 can observe in the range from about 0.2 μm to 1.7 μm

5.
Angular resolution
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In scientific analysis, in general, the term resolution is used to describe the precision with which any instrument measures and records any variable in the specimen or sample under study. The imaging systems resolution can be limited either by aberration or by diffraction causing blurring of the image and these two phenomena have different origins and are unrelated. Aberrations can be explained by geometrical optics and can in principle be solved by increasing the optical quality —, on the other hand, diffraction comes from the wave nature of light and is determined by the finite aperture of the optical elements. The lens circular aperture is analogous to a version of the single-slit experiment. The interplay between diffraction and aberration can be characterised by the point spread function, the narrower the aperture of a lens the more likely the PSF is dominated by diffraction. If the distance is greater, the two points are well resolved and if it is smaller, they are regarded as not resolved, Rayleigh defended this criteria on sources of equal strength. Considering diffraction through an aperture, this translates into, θ =1.220 λ D where θ is the angular resolution, λ is the wavelength of light. The factor 1.220 is derived from a calculation of the position of the first dark circular ring surrounding the central Airy disc of the diffraction pattern and this number is more precisely 1.21966989. The first zero of the order-one Bessel function of the first kind J1 divided by π, the formal Rayleigh criterion is close to the empirical resolution limit found earlier by the English astronomer W. R. Dawes who tested human observers on close binary stars of equal brightness. The result, θ =4. 3% dip, modern image processing techniques including deconvolution of the point spread function allow resolution of binaries with even less angular separation. The angular resolution may be converted into a resolution, Δℓ. For a microscope, that distance is close to the length f of the objective. For this case, the Rayleigh criterion reads, Δ ℓ =1.220 f λ D. This is the size, in the plane, of smallest object that the lens can resolve. The size is proportional to wavelength, λ, and thus, for example and this result is related to the Fourier properties of a lens.220 f λ D =1.22 λ ⋅. Since this is the radius of the Airy disk, the resolution is better estimated by the diameter,2.44 λ ⋅ Point-like sources separated by a smaller than the angular resolution cannot be resolved. A single optical telescope may have a resolution less than one arcsecond. The angular resolution R of a telescope can usually be approximated by R = λ D where λ is the wavelength of the observed radiation, the Resulting R is in radians

6.
Radio telescope
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A radio telescope is a specialized antenna and radio receiver used to receive radio waves from astronomical radio sources in the sky in radio astronomy. Radio telescopes are typically large parabolic antennas similar to those employed in tracking and communicating with satellites and they may be used singly, or linked together electronically in an array. Unlike optical telescopes, radio telescopes can be used in the daytime as well as at night, Radio waves from space were first detected by engineer Karl Guthe Jansky in 1932 at Bell Telephone Laboratories in Holmdel, New Jersey using an antenna built to study noise in radio receivers. The first purpose-built radio telescope was a 9-meter parabolic dish constructed by radio amateur Grote Reber in his yard in Wheaton. The sky survey he did with it is considered the beginning of the field of radio astronomy. The first radio antenna used to identify an astronomical radio source was one built by Karl Guthe Jansky, Jansky was assigned the job of identifying sources of static that might interfere with radio telephone service. Janskys antenna was an array of dipoles and reflectors designed to receive short wave radio signals at a frequency of 20.5 MHz and it was mounted on a turntable that allowed it to rotate in any direction, earning it the name Janskys merry-go-round. It had a diameter of approximately 100 ft and stood 20 ft tall, by rotating the antenna, the direction of the received interfering radio source could be pinpointed. A small shed to the side of the antenna housed an analog recording system. Jansky finally determined that the faint hiss repeated on a cycle of 23 hours and 56 minutes and this period is the length of an astronomical sidereal day, the time it takes any fixed object located on the celestial sphere to come back to the same location in the sky. An amateur radio operator, Grote Reber, was one of the pioneers of what became known as radio astronomy and he built the first parabolic dish radio telescope, a 9 metres in diameter) in his back yard in Wheaton, Illinois in 1937. The range of frequencies in the spectrum that makes up the radio spectrum is very large. This means that the types of antennas that are used as radio telescopes vary widely in design, size, at wavelengths of 30 meters to 3 meters, they are generally either directional antenna arrays similar to TV antennas or large stationary reflectors with moveable focal points. Since the wavelengths being observed with these types of antennas are so long, at shorter wavelengths parabolic dish antennas predominate. The angular resolution of an antenna is determined by the ratio of the diameter of the dish to the wavelength of the radio waves being observed. This dictates the size a radio telescope needs for a useful resolution. Radio telescopes that operate at wavelengths of 3 meters to 30 cm are usually well over 100 meters in diameter, telescopes working at wavelengths shorter than 30 cm range in size from 3 to 90 meters in diameter. The Wilkinson Microwave Anisotropy Probe mapped the Cosmic microwave background radiation in 5 different frequency bands, centered on 23 GHz,33 GHz,41 GHz,61 GHz, the worlds largest filled-aperture radio telescope is the Five hundred meter Aperture Spherical Telescope completed in 2016 by China

7.
Karl G. Jansky Very Large Array
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The Karl G. Jansky Very Large Array is a radio astronomy observatory located on the Plains of San Agustin, between the towns of Magdalena and Datil, some 50 miles west of Socorro, New Mexico. It comprises 27 25-meter radio telescopes in a Y-shaped array and all the equipment, instrumentation, the VLA stands at an elevation of 6970 ft above sea level. It is a component of the National Radio Astronomy Observatory, the radio telescope comprises 27 independent antennae, each of which has a dish diameter of 25 meters and weighs 209 metric tons. The antennae are distributed along the three arms of a track, shaped in a wye -configuration, using the rail tracks that follow each of these arms—and that, at one point, intersect with U. S. The angular resolution that can be reached is between 0.2 and 0.004 arcseconds, there are four commonly used configurations, designated A through D. The observatory normally cycles through all the possible configurations every 16 months. Moves to smaller configurations are done in two stages, first shortening the east and west arms and later shortening the north arm and this allows for a short period of improved imaging of extremely northerly or southerly sources. The frequency coverage is 74 MHz to 50 GHz, the Array Operations Center for the VLA is located on the campus of the New Mexico Institute of Mining and Technology in Socorro, New Mexico. In 2011, a long upgrade project had resulted in the VLA expanding its technical capacities by factors of as much as 8,000. The 1970s era electronics were replaced with state-of-the-art equipment, on March 31,2012, the VLA was officially renamed in a ceremony inside the Antenna Assembly Building. In 1989 the VLA was used to receive communications from the Voyager 2 spacecraft as it flew by Neptune. It is not, despite depictions in culture such as the movie Contact. It has been used to carry out several surveys of radio sources, including the NRAO VLA Sky Survey. The driving force for the development of the VLA was David S. Heeschen and he is noted as having sustained and guided the development of the best radio astronomy observatory in the world for sixteen years. Congressional approval for the VLA project was given in August 1972, the first antenna was put into place in September 1975 and the complex was formally inaugurated in 1980, after a total investment of USD $78.5 million. It was the largest configuration of radio telescopes in the world, with a view to upgrading the venerable 1970s technology with which the VLA was built, the VLA has evolved into the Expanded Very Large Array. The upgrade has enhanced the sensitivity, frequency range. A second phase of this upgrade may add up to eight additional dishes in other parts of the state of New Mexico, up to 300 km away, in 1995 and 1996, the VLA was used for following up the Wow. signal from the SETI project

8.
Very-long-baseline interferometry
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Very-long-baseline interferometry is a type of astronomical interferometry used in radio astronomy. In VLBI a signal from a radio source, such as a quasar, is collected at multiple radio telescopes on Earth. The distance between the telescopes is then calculated using the time difference between the arrivals of the radio signal at different telescopes. This allows observations of an object that are made simultaneously by many radio telescopes to be combined, Data received at each antenna in the array include arrival times from a local atomic clock, such as a hydrogen maser. At a later time, the data are correlated with data from other antennas that recorded the same radio signal, the resolution achievable using interferometry is proportional to the observing frequency. VLBI is most well known for imaging distant cosmic radio sources, spacecraft tracking, using VLBI in this manner requires large numbers of time difference measurements from distant sources observed with a global network of antennas over a period of time. Some of the results derived from VLBI include, High resolution radio imaging of cosmic radio sources. Variations in the Earths orientation and length of day, the most sensitive VLBI array in the world is the European VLBI Network. This is an array which brings together the largest European radiotelescopes for typically week-long sessions. The combination of the EVN and VLBA is known as Global VLBI, VLBI has traditionally operated by recording the signal at each telescope on magnetic tapes or disks, and shipping those to the correlation center for replay. Recently, it has become possible to connect VLBI radio telescopes in close to real-time, while employing the local time references of the VLBI technique. The image to the shows the first science produced by the European VLBI Network using e-VLBI. The data from 6 telescopes were processed in time at the European Data Processing centre at JIVE. The Netherlands Academic Research Network SURFnet provides 6 x 1 Gbit/s connectivity between JIVE and the GEANT2 network, in the quest for even greater angular resolution, dedicated VLBI satellites have been placed in Earth orbit to provide greatly extended baselines. Experiments incorporating such space-borne array elements are termed Space Very Long Baseline Interferometry, another space VLBI mission, Spektr-R, was launched in July 2011. In VLBI interferometry, the digitized antenna data are recorded at each of the telescopes. The antenna signal is sampled with a precise and stable atomic clock that is additionally locked onto a GPS time standard. Alongside the astronomical data samples, the output of this clock is recorded on the tape/disk media, the recorded media are then transported to a central location

9.
Submillimeter Array
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The Submillimeter Array consists of eight 6-meter diameter radio telescopes arranged as an interferometer for submillimeter wavelength observations. All three of these observatories are located at Mauna Kea Observatory on Mauna Kea, Hawaii, the baseline lengths presently in use range from 16 to 508 meters, and up to 783 meters for eSMA operations. Although the array is capable of operating both day and night, most of the observations take place at nighttime when the phase stability is best. The SMA is jointly operated by the Smithsonian Astrophysical Observatory and the Academia Sinica Institute of Astronomy, the SMA is a multi-purpose instrument which can be used to observe diverse celestial phenomena. The SMA excels at observations of dust and gas with only a few tens of kelvins above absolute zero. Commonly observed classes of objects include star-forming molecular clouds in our own and other galaxies, highly redshifted galaxies, evolved stars, occasionally, bodies in the Solar System, such as planets, asteroids, comets and moons, are observed. The SMA has been used to discover that Pluto is 10 K cooler than expected and it was the first radio telescope to resolve Pluto and Charon as separate objects. The SMA is a part of the Event Horizon Telescope, which observes nearby supermassive black holes with a resolution comparable to the size of the objects event horizon. 4″N 155°28′40. 7″W

10.
LOFAR
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LOFAR consists of a vast array of omnidirectional antennas using a new concept in which the signals from the separate antennas are not combined in real time as they are in most array antennas. The electronic signals from the antennas are digitized, transported to a digital processor. The project is based on an array of radio telescopes using about 20,000 small antennas concentrated in at least 48 stations. Forty of these stations are distributed across the Netherlands and were funded by ASTRON, the five stations in Germany, and one each in Great Britain, France and Sweden, were funded by these countries. Further stations may also be built in other European countries, the total effective collecting area is approximately 300,000 square meters, depending on frequency and antenna configuration. The data processing is performed by a Blue Gene/P supercomputer situated in the Netherlands at the University of Groningen, LOFAR is also a technology precursor for the Square Kilometre Array. LOFAR was conceived as an effort to force a breakthrough in sensitivity for astronomical observations at radio-frequencies below 250 MHz. Astronomical radio interferometers usually consist either of arrays of parabolic dishes, the direction of observation of the stations is chosen electronically by phase delays between the antennas. LOFAR can observe in several directions simultaneously, as long as their data rate remains under its cap. This in principle allows a multi-user operation, LOFAR makes observations in the 10 MHz to 240 MHz frequency range with two types of antennas, Low Band Antenna and High Band Antenna, optimized for 10-80 MHz and 120-240 MHz respectively. The electric signals from the LOFAR stations are digitised, transported to a digital processor. Therefore, LOFAR is a software telescope, the cost is dominated by the cost of electronics and will therefore mostly follow Moores law, becoming cheaper with time and allowing increasingly large telescopes to be built. The antennas are simple enough, but there are about 20,000 in the LOFAR array, to make radio surveys of the sky with adequate resolution, the antennas are arranged in clusters that are spread out over an area of more than 1000 km in diameter. The LOFAR stations in the Netherlands reach baselines of about 100 km, LOFAR currently receives data from 24 core stations,14 remote stations in The Netherlands and 8 international stations. Each of the core and remote stations has 48 HBAs and 96 LBAs, international stations have 96 LBAs and 96 HBAs and a total of 96 digital Receiver Units. This installation functions as a VHF receiver either in stand-alone mode or part of a radar system together with EISCAT transmitter in Tromsø. Data transport requirements are in the range of several gigabits per second per station, the data from LOFAR is stored in the LOFAR long-term archive. LOFAR is the most sensitive radio observatory at its low observing frequencies, until the generation of large array radio telescope

This orange blob shows the nearby star Betelgeuse, as seen by the Atacama Large Millimeter/submillimeter Array (ALMA). This is the first time that ALMA has ever observed the surface of a star and this first attempt has resulted in the highest-resolution image of Betelgeuse available.

Coronagraphic image of AB Pictoris showing a companion (bottom left), which is either a brown dwarf or a massive planet. The data was obtained on 16 March 2003 with NACO on the VLT, using a 1.4 arcsec occulting mask on top of AB Pictoris.

The 'Zernikeborg' building, which houses the University of Groningen's computing center

At low radio frequencies the sky is dominated by small bright sources (shown is a 151 MHz map of the region: 140° to 180° Galactic longitude; -5° to 5° Galactic latitude). LOFAR will have sufficient fidelity and sensitivity to see faint structure between these bright sources because of the very large number of array elements.

Materials with higher emissivity appear to be hotter. In this thermal image, the ceramic cylinder appears to be hotter than its cubic container (made of silicon carbide), while in fact they have the same temperature.

Aperture synthesis or synthesis imaging is a type of interferometry that mixes signals from a collection of telescopes …

Most aperture synthesis interferometers use the rotation of the Earth to increase the number of baseline orientations included in an observation. In this example with the Earth represented as a grey sphere, the baseline between telescope A and telescope B changes angle with time as viewed from the radio source as the Earth rotates. Taking data at different times thus provides measurements with different telescope separations.

Angular resolution or spatial resolution describes the ability of any image-forming device such as an optical or radio …

Airy diffraction patterns generated by light from two points passing through a circular aperture, such as the pupil of the eye. Points far apart (top) or meeting the Rayleigh criterion (middle) can be distinguished. Points closer than the Rayleigh criterion (bottom) are difficult to distinguish.

Log-log plot of aperture diameter vs angular resolution at the diffraction limit for various light wavelengths compared with various astronomical instruments. For example, the blue star shows that the Hubble Space Telescope is almost diffraction-limited in the visible spectrum at 0.1 arcsecs, whereas the red circle shows that the human eye should have a resolving power of 20 arcsecs in theory, though normally only 60 arcsecs.

Aperture Masking Interferometry is a form of speckle interferometry, that allows diffraction limited imaging from …

a) shows a simple experiment using an aperture mask in a re-imaged aperture plane. b) and c) show diagrams of aperture masks which were placed in front of the secondary mirror of the Keck telescope by Peter Tuthill and collaborators. The solid black shapes represent the subapertures (holes in the mask). A projection of the layout of the Keck primary mirror segments is overlaid.